Telescope Array Experiment1

Telescope Array Experiment
H. Kawaia, S. Yoshidaa , H. Yoshiib , K. Tanakac , F. Cohend , M. Fukushimad , N. Hayashidad,
K. Hiyamad , D. Ikedad, E. Kidod , Y. Kondod , T. Nonakad , M. Ohnishid , H. Ohokad , S. Ozawad,
H. Sagawad, N. Sakuraid, T. Shibatad , H. Shimodairad , M. Takedad , A. Taketad , M. Takitad ,H. Tokunod, R. Toriid , S. Udod , Y. Yamakawad, H. Fujiie , T. Matsudae , M. Tanakae , H. Yamaokae,
K. Hibinof , T. Bennog , K. Dourag , M. Chikawag , T. Nakamurah, M. Teshimai , K. Kadotaj,
Y. Uchihorik , K. Hayashil, Y. Hayashil, S. Kawakamil, T. Matsuyamal , M. Minaminol , S. Ogiol∗ ,A. Ohshimal , T. Okudal , N. Shimizul , H. Tanakal , D.R. Bergmanm, G. Hughesm , S. Strattonm ,G.B. Thomsonm , A. Endon , N. Inouen , S. Kawanan, Y. Wadan , K. Kasaharao, R. Azumap , T. Iguchip ,F. Kakimotop , S. Machidap , K. Misumip , Y. Muranop , Y. Tamedap , Y. Tsunesadap , J. Chibaq ,K. Miyataq , T. Abu-Zayyadr, J.W. Belzr , R. Cadyr , Z. Caor , P. Huentemeyerr , C.C.H. Juir ,K. Martensr , J.N. Matthewsr , M. Mostofar , J.D. Smithr , P. Sokolskyr, R.W. Springerr , J.R. Thomasr ,S.B. Thomasr , L.R. Wiencker , T. Doyles , M.J. Taylors , V.B. Wickwars , T.D. Wilkersons ,K. Hashimotot , K. Hondat , K. Ikutat , T. Ishiit , T. Kanbet , and T. TomidataChiba University, 1-33 Yayoi-cho, Inage-ku, Chiba, 263-8522, Japanb

Ehime University, 2-5 Bunkyo-cho, Matsuyama, Ehime, 790-8577, Japanc

Hiroshinma City University, 3-4-1 Ozuka-Higashi, Asa-Minami-ku, Hiroshima, 731-3194, Japandef

ICRR, University of Tokyo, 5-1-5 Kashiwanoha, Kashiwa, Chiba, 277-8582, Japan

Institute of Particle and Nuclear Studies, KEK, 1-1 Oho, Tsukuba, Ibaraki, 305-0801, Japan

Kanagawa University, 3-27-1 Rokkakubayashi, Kanagawa-ku, Yokohama, Kanagawa, 221-8686, Japang

Kinki University, 3-4-1 Kowakae, Higashi-Osaka, Osaka, 577-8502, Japanh

Kochi University, 2-5-1 Akebono-cho, Kochi, 780-8520, Japani

Max-Planck-Institute for Physics, Foehringer Ring 6, 80805 München, GermanyjMusashi Institute of Technology, 1-28-1 Tamazutsumi, Setagawa-ku, Tokyo, 158-8558, JapanklNational Institute of Radiological Sciences, 4-9-1 Anagawa Inage-ku, Chiba, 263-8555, Japan

Osaka City University, 3-3-138 Sugimoto-cho, Sumiyoshi-ku, Osaka, 558-8585, Japanm

Rutgers University, 136 Freilinghuysen Road, Piscataway, NJ 08854, USAn

Saitama University, 255 Shimo-Okubo, Sakura-ku, Saitama, 338-8570, Japano

Shibaura Institute of Technology, 307 Fukasaku, Minuma-ku, Saitama, 337-8570, Japanp

Tokyo Institute of Technology, 2-12-1 Ookayama, Meguro-ku, Tokyo, 152-8550, Japanq

Tokyo University of Science, 2641 Yamazaki, Noda, Chiba, 278-8510, Japanr

University of Utah, 115 S 1400 E, Salt Lake City, UT 84112, USAs

Utah State University, Logan, UT 84322, USAt

Yamanashi University, 4-4-37 Takeda, Kofu, Yamanashi, 400-8510, JapanThe TA observatory is a hybrid detector system consisting of both a surface detector array as well as a setof fluorescence detectors, and the observatory will measure the energy spectrum, anisotropy and composition ofultra-high energy cosmic rays. The Surface Detectors are being deployed and the array should be complete by
the end of February, 2007. We will soon be collecting hybrid data at the Telecope Array.

verted to electro–magnetic component in the atmosphere, and because the lateral particle density
distribution of electrons has a smaller dependenceon hadron interaction models[6] and on primary
species than that of muons.

1. Introduction
The Telescope Array (TA) Project[1] is designed to study cosmic rays with energies above
1019 eV. The TA observatory will measure theenergy spectrum, arrival direction distribution(anisotropy) and composition of Ultra-High Energy Cosmic Rays (UHECRs). It will look for thesuper–GZK events[2] and the event clusters[3] observed by AGASA (the Akeno Giant Air Shower
Array) and will resolve the discrepancy between
the AGASA results and those of HiRes (the High
Resolution Fly’s Eye)[4].
The TA observatory is a hybrid detector system consisting of both a Surface Detector (SD)array (ala AGASA) as well as a set of Fluorescence Detectors (FD) (ala HiRes). A map of theobservatory is shown in Figure 1, i.e., it consistsof 576 scintillation Surface Detectors (SDs) whichmeasure the distribution of charged particles at
the Earth’s surface and three Fluorescence Detector (FD) stations which observe the night sky
above the SD array[5]. The SD array covers an
area nine times greater than the AGASA detection area. We can observe both the longitudinaldevelopment and the lateral distribution of air
showers simultaneously with the hybrid system.
Thus, we will use the measured parameters forcomplementary calibrations and can improve theenergy and the angular resolutions of each component of the hybrid system.
The TA observatory is located in Millard
County, Utah, USA(39.1◦ N, 112.9◦W). The site
is about 1400 m above sea level. As with the
AGASA and HiRes experiments, the TA observatory is located in the northern hemisphere. This
allows the TA to directly compare with AGASA
and HiRes measurement and to compare withsouthern experiments to search for a potential
north–south asymmetry in energey spectra orsource distribution.
We chose scintillation detectors for the SD array of TA experiment, because scintillators havea smaller systematic error in the primary energy
determination when compared to a water tank
SD array. This is because, for UHECRs, about
90% of the energy of the primary particle is con-Figure 1. Detector arrangement of the TA SDs(diamond) and the FD stations (pentagon).

2. Surface Detector
The surface detectors are distributed on a 1.2
km square grid and the effective detection area of
the whole SD array will be 762 km2 . Each SD(Figure 2) has an assembled weight of 250 kg andconsists of a power supply, two layers of scintillation detector and electronics. Power is generated
by a 120 W solar panel and stored in a sealed
lead-acid battery. The system has the capacity
for one week’s operation in complete darkness.
Each scintillation detector layer is made of extruded plastic scintillator 1.2 cm thick and withan area of 3 m2 . The photo multiplier tube
(PMT, Electrontube 91245A with 9/8 inch diameter) is connected to the scintillator via 96
wave length shifting fibers (WLSFs, 2mm diameter). The fibers rest in grooves in the scintillator∗ Corresponding Author: S. Ogio, sogio@sci.osaka-cu.ac.jp

2Telescope Array Experiment

Figure 2. A surface detector deployed in the field.

aligned parallel to the long side of the scintillator
panels and have 2 cm separation. The averaged
number of photo electrons induced by a penetrating Minimum Ionizing Particle (MIP) is 24. The
PMT is set to a gain of about 4 × 106 .
The analog part of the SD electronics[7] has anlow–pass filter with cut–off frequency of 9.7 MHz
for each PMT, and the output is readout by a
50 MHz 12 bit FADC. The ADC delivers its data
to a buffer which has sufficient capacity to store
data for 16,000 events at 5.12 µs/event.
The event trigger for the SD array is implemented in a Field Programable Gate Array(FPGA). At present, we intend to use the threefold coincidence of nearby detectors with morethan three particle hits. In this case, the trigger recognition time is less than 10 ms. Using
the above triggering conditions, we calculated thetriggering efficiencies for UHECRs which haveincident zenith angles less than 45◦ . The efficiency for primary protons is 35% at an energy of1018 eV and 100% at 1019 eV. The estimated energy resolutions from analyzing simulated events
is about 27% at 1019 eV and 19% at 1020 eV,
which comparable in accuracy to AGASA.
3. FD Station, Telescope and Camera
The Telescope Array consists of three Fluorescence Detector (FD) stations. The stations are3located on a triangle with about 35 km separation and consists of 12-14 telescopes. Each station is centered on the central laser facility in the
middle of the array and views 3◦ -33◦ in elevationand ∼108◦ in azimuth, and thus each one views
nearly the entire array at high energies.
We have completed installation of telescopes
at the first FD station in the south–east cornerof the array, and we are currently installing thetelescopes at the second station in the south-west
corner. The third station near the north–west
corner is currently under construction. The site
will be complete by the end of 2006 and we plan
to move some of the HiRes–I telescopes to thisstation.
The telescopes (Figure 3) at the first two sites
have a combined spherical mirror with a diameterof 3.3 m, a focal length of 3.0 m and a spot size of
30 mm on the focal plane. The pointing accuracyof the telescopes is 0.07◦ .
Each telescope has a fluorescence light camera which consists of 256 PMTs (HAMAMATSU
R9508). The sensitive area of the camera is 1 m
× 1 m, which corresponds to the field of view of
15◦ in elevation times 18◦ in azimuth. Thus the
pixel size of the camera is nearly equal to 1◦ .
The high voltage to the PMT is negative andis applied to the PMTs individually with a computer controlled power supply system. This allows us to match the gains of all the PMTs and tocompensate for aging effects of the PMTs. Moreover, the PMTs are DC–coupled so that the direct measurement of the DC–anode current ofnight sky background at the front-end electronics is possible.
In each camera three PMTs among the 256
are absolutely calibrated with a standard lightsource called “CRAYS”[8], which exposes PMTs
to Rayleigh–scattered photons out of a calibrated
N2 laser (337.1 nm, 300 µJ) in a N2 –filled chamber. Moreover, the gain of each calibrated PMT
is monitored through the constant measurements
using a YAP light pulsar attached on the photo
cathode. The YAP pulsar[9] consists of 241 Am
and a YAlO3 :Ce scintillator, and its dimension
are 4 mm in diameter and 2 mm in thickness.
The gains of other PMTs are calibrated relative
to the absolutely calibrated PMTs by comparing

4

S. Ogio

Figure 4. A block diagram of the electronic andthe DAQ system in a FD station.

Figure 3. Telescopes and cameras in the first TA
FD station. Every segment mirror of the upper
mirror has a cover to avoid dust.

output signals of PMTs when they are exposed
to a Xe flasher installed at the center of a mirror.
The accuracy of absolute and relative gain determinations is 7% and 2.5%, respectively. Finally,
the PMT–gain is set at 8 × 104 and the gain of
the pre–amplifiers at the PMTs is 50.
4. FD Electronics
The schematic of the electronics and the dataacquisition (DAQ) system of a FD station isshown in Figure 4. The FD electronics[10] consists of the following three main componenta: the
Signal Digitizer/Finder (SDF), the Track Finder(TF), and the Central Trigger Distributer (CTD).
Each component fits on a single width VME–9U
module.
Each SDF board processes the signals of one
camera column, i.e., there are 16 channels per
board. Each channel has a FADC to readout

the pre–amplifier, a buffer memory, and a signal
finder logic. The analog input is split into a low
gain (1/16) sum per 16 channels. The sum channel also has a FADC and a buffer memory. With
the sum channels, the effective dynamic range
and the sampling frequency of SDF is equivalent
to 16 bits at 10MHz. The signal finder logic measures the significance of the excess photons abovenight sky background in every 12.8 µs window,
which we call “frames”. In a frame, the sum of
ADC values of the first N samples are comparedwith the fluctuation of DC current which is calculated from the last 125 frames. This is processed
for four different sum window size N = 16, 32, 64
and 128, and the recognitions are repeated withshifting the sum window by 1 sampling time until
the window reaches the end of the frame. In theseprocesses, when the net count which is the averaged DC current from the sum exceeds a threshold, this channel is recognized as a “hit” channel
in the frame. The threshold of the significancelevel is programmable. The time for the signal
finding in a frame is 1.8 µs.
Each TF processes the hit patterns of one camera in every frame. The TF recognizes 5–fold hit
pixel patterns as a track and 4–fold hit on theedge of the camera as a partial track, and sends

Telescope Array Experiment
the track recognition result to the CTD. The processing time for single track finding routine is 5.4µs.
One CTD processes the track recognition results of whole the camera in a FD station. Whenone or more TFs find tracks, CTD distributes
trigger signals to start data recording procedures
for all the SDFs and the TFs. Moreover, wheneach of two neighboring TFs finds a partial track,
the CTD also triggers the DAQ system. The total
triggering time including SDF and TF processing
is 9.8 µs. In addition, the CTD takes care of otherimportant functions – such as tracking the GPS
clock, external trigger input, synchnonization of
the system clock for all the electronics of the DAQ
system in a station.
The triggering efficiencies of the FD stations forstereo coincidence are shown in Figure 5. In theenergy range of 1019 eV and above the detection
area of stereo FD almost completely covers the
SD array.

Figure 5. The core locations of the incident
cosmic rays (green open circle) and the stereo–
triggered showers (red filled circle) for primary
protons of an incident zenith angle of 30◦ . The
primary energies are 1019 eV (left) and 1020 eV(right). Each of the blue filled triangles indicatesa location of a FD station.

The DAQ computers are twelve Mirror PCs,a Station PC for CTD, a Slow Control PC, a
Data Storage PC, an Analysis PC and a GUI
PC. We use the Linux operating system for every

5
PC and use a Network Shared Memory (NSM) library which was developed for the DAQ system of
the Belle experiment at KEK. This library provides PC to PC communication services similar
to “signal” and “shared–memory” on UNIX. The
data size per event is 250 kB per camera. Consequently, the amount of data per night is 10 GB
per station, when the triggering rate is 1.2 Hz
per station. We will expect to regulate the SDF
threshold level to keep the trigger rate at about 1
Hz per station, because the expected air shower
event rate is 0.005 Hz. However, the DAQ system
has the potential to work with a maximum rateof 20 Hz.
5. Atmospheric Monitoring and Calibrations
In order to monitor the condition of the atmosphere, we use two monitoring systems. The
first is a steerable laser (used in bi-static LIDARmode) which is located at the Central Laser Facility (CLF). This provides a “test beam” to investigate properties of the atmosphere and theFDs. The CLF is located at the center of the array and is equidistant from each FD station (20.85km). The CLF laser is an energy tripled Nd:YAG
laser(λ = 355 nm) with an output of 5 mJ. Thus,
at a full energy, the number of photons out of abeam is 3.8 × 1011 photons/m which is approximately equivalent to fluorescence photons emitted by showers with energy of 1020 eV,i.e., about
4 × 1011 photons/m.
The second atmospheric monitoring system ais standard LIDAR system[11]. It consists of alaser and light detectors to measure the intensity
and the time profiles of back scattered photons.
One such station will be installed at each of the
FD stations. The LIDAR at the first FD station
has a Nd:YAG laser (355 nm) with an output of
5 mJ and a photon detector which consists of a
f/10 Cassegrain–telescope with a diameter of 30.5
cm and a PMT with a diameter of 19 mm.
For an end–to–end calibration of FDs, we havea plan to use a linear electron accelerator (40MeV/electron) in a trailer. We will shoot theaccelerator vertically in the FD’s field of view at
a distance of 100 meters. The current from the6accelerator is 109 electrons per shot and the repetition frequency is 1 Hz. The amount of energy
in one shot of the LINAC is 4 × 1016 eV, whichis equivalent to 1020 eV at a distance of 10 km.
If we can use this “test beam” calibration, we
will reduce the energy determination uncertainty
of 20% to 14% including the uncertainty due toatmospheric monitoring.
6. Conclusion
The Telescope Array Project is under construction in central Utah, USA. The Surface Detectors are being deployed and the array should becomplete by the end of February, 2007. The firstof three Fluorescence Detector stations is nearlycomplete. The building for the second is complete
and the site should be operational in 3/2007. The
Third site is under construction and should be
operational in 5/2007. We will soon be collecting
hybrid data at the Telecope Array.
7. Acknowledgments
The Telescope Array collaboration wishes toacknowledge the support of the the Japanese government through a Grant-in-Aid for Scientific Research (Kakenhi) on the Priority Area ”The Origin of the Highest Energy Cosmic Rays” as wellas the U.S. National Science Foundation (NSF)
through awards PHY-0307098 and PHY-0601915(University of Utah) and PHY-0305516 (Rutgers
University). The Dr. Ezekiel R. and Edna Wattis
Dumke Foundation, The Willard L. Eccles Foundation and The George S. and Dolores Dore Eccles Foundation all helped with generous donations. The State of Utah supported the projectthrough its Economic Development Board, andthe University of Utah supported us throughthe Office of the Vice President for Research.
We gratefully acknowledge the contributions fromthe technical staffs of our home institutions andthe Utah Center for High Performance Computing. The TA could not be built without strong
support from the agencies that administer thelands and airspace above our project area. We
thank the State of Utah School and Institutional
Trust Lands Administration, the federal Bureau

S. Ogioof Land Management, and the United States Air
Force. We also wish to thank the people and officials of Millard County, Utah, for their steadfast
support of our experiment.
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